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EHowever the amplitude of variability in NGC€43395 is comparable to that of Νας0521 observed with Ginga (Llavashicda ct al 1998).
However the amplitude of variability in NGC4395 is comparable to that of NGC4051 observed with Ginga (Hayashida et al 1998).
rapidly.
rapidly.
While the knots of 330 are broadened by ! (Meaburn Lóppez 1996). the echelle spectra of Pollacco et al. (
While the knots of 30 are broadened by $<$ $^{-1}$ (Meaburn Lóppez 1996), the echelle spectra of Pollacco et al. (
1992) yielded velocity FWHMs for the knot of + from the [N 1] lines and + from the [O 111] A5007 line.
1992) yielded velocity FWHMs for the knot of $^{-1}$ from the [N ] lines and $^{-1}$ from the [O ] $\lambda$ 5007 line.
Sakurai’s Object (V4334 Ser) was discovered in February 1996. and was initially thought to be a slow nova (Nakano et al.
Sakurai's Object (V4334 Sgr) was discovered in February 1996, and was initially thought to be a slow nova (Nakano et al.
1996).
1996).
However. its spectrum showed a rapid cooling. hydrogen abundance decline and an enhancement of s-process elements. over just a few months after its discovery. and it is thought to have undergone a tinal helium flash (Duerbeck Benetti 1996).
However, its spectrum showed a rapid cooling, hydrogen abundance decline and an enhancement of s-process elements, over just a few months after its discovery, and it is thought to have undergone a final helium flash (Duerbeck Benetti 1996).
The spectrum of Sakurai's object at its maximum brightness is very similar to the spectrum of V605 Aq! obtained at its maximum in 1919. with the appearance of a cool RCB star (Lundmark 1921).
The spectrum of Sakurai's object at its maximum brightness is very similar to the spectrum of V605 Aql obtained at its maximum in 1919, with the appearance of a cool RCB star (Lundmark 1921).
Because of this. V605 Aqgl is often described as an older twin of Sakurai's Object.
Because of this, V605 Aql is often described as an older twin of Sakurai's Object.
The He:C:O mass ratios at the surface of Sakurai’s object were found to be 85:5:3 in July 1996. compared to 54:40:5 for the central star of 558 (Clayton et al.
The He:C:O mass ratios at the surface of Sakurai's object were found to be 85:5:3 in July 1996, compared to 54:40:5 for the central star of 58 (Clayton et al.
2006).
2006).
Our derived nebular abundances for 5587s knot correspond to He:C:O mass ratios of [8:2:41.
Our derived nebular abundances for 58's knot correspond to He:C:O mass ratios of 18:2:41.
Freshly ionised material has been detected around Sakurai’s Object. indicating that a knot or knots similar to those seen in 558 may be forming.
Freshly ionised material has been detected around Sakurai's Object, indicating that a knot or knots similar to those seen in 58 may be forming.
The freshly ionised material is found to be expanding at + (Kerber et al.
The freshly ionised material is found to be expanding at $^{-1}$ (Kerber et al.
2002).
2002).
At this expansion rate. its density will decline rapidly and it seems unlikely that any knot formed could survive to be seen as long after the event as the knots of 558 and A330.
At this expansion rate, its density will decline rapidly and it seems unlikely that any knot formed could survive to be seen as long after the event as the knots of 58 and 30.
The physical conditions and chemical abundances in the freshly ionised ejecta have yet to be determined.
The physical conditions and chemical abundances in the freshly ionised ejecta have yet to be determined.
Possible links between planetary nebulae showing high adfs and classical novae have been discussed by Wesson et al. (
Possible links between planetary nebulae showing high adfs and classical novae have been discussed by Wesson et al. (
2003) and Liu et al. (
2003) and Liu et al. (
2006).
2006).
In several cases. nova shells have been found to show very strong recombination line emission and evidence for very low temperatures.
In several cases, nova shells have been found to show very strong recombination line emission and evidence for very low temperatures.
One good example is the shell surrounding DQ Her. which was found by Williams et al. (
One good example is the shell surrounding DQ Her, which was found by Williams et al. (
1978) to have a Balmer jump temperature of —-500K. The C/O number ratio in the ejecta of DQ Her is 0.36 — quite similar to that seen in the knots of Abell 30.
1978) to have a Balmer jump temperature of $\sim$ 500K. The C/O number ratio in the ejecta of DQ Her is 0.36 – quite similar to that seen in the knots of Abell 30.
CP Pup shows similar spectral features with very strong recombination lines and a Balmer jump temperature of 5060 KK. but in its case no carbon recombination lines are seen. implying a lower C/O ratio (Williams 1982).
CP Pup shows similar spectral features with very strong recombination lines and a Balmer jump temperature of $\sim$ K, but in its case no carbon recombination lines are seen, implying a lower C/O ratio (Williams 1982).
We have included in Table 8. a summary of the elemental mass fractions measured for two old nova shells and for three "neon novae'.
We have included in Table \ref{comparisons} a summary of the elemental mass fractions measured for two old nova shells and for three `neon novae'.
All five show C/O ratios of less than unity but the neon novae also show neon mass fractions that are comparable to the large values (€0.1—0.4) that we have measured for the knots of Α 58 and A 30.
All five show C/O ratios of less than unity but the neon novae also show neon mass fractions that are comparable to the large values (0.1–0.4) that we have measured for the knots of A 58 and A 30.
Models for neon novae (e.g. Starrtield et al.
Models for neon novae (e.g. Starrfield et al.
1986: Politano et al.
1986; Politano et al.
1995) invoke the usual thermonuclear runaway on the surface of a white dwarf following mass transfer from a low mass companion. with the high neon abundances resulting from the fact that the runaway occurs on the surface of a high-mass 1.0-1.35 . ) O-Ne-Mz white dwarf. some of whose subsurface material is mixed to the surface and ejected during the nova event.
1995) invoke the usual thermonuclear runaway on the surface of a white dwarf following mass transfer from a low mass companion, with the high neon abundances resulting from the fact that the runaway occurs on the surface of a high-mass (1.0-1.35 $_\odot$ ) O-Ne-Mg white dwarf, some of whose subsurface material is mixed to the surface and ejected during the nova event.
This suggests that the central stars of A 58 and A 30 might have high mass O-Ne-Mg cores. some of whose material may been brought to localised parts of the surface by the event that led to the ejection of the observed knots.
This suggests that the central stars of A 58 and A 30 might have high mass O-Ne-Mg cores, some of whose material may been brought to localised parts of the surface by the event that led to the ejection of the observed knots.
If entropy barriers prevent such a scenario during a VLTP event then an alternative scenario could involve the transfer of mass from a companion on to localised parts of a massive white dwarf surface. leading to a thermonuclear runaway and the excavation and ejection of material from the Ne-Mg region of the white dwarf. whose thin surface layer still has C/O > |.
If entropy barriers prevent such a scenario during a VLTP event then an alternative scenario could involve the transfer of mass from a companion on to localised parts of a massive white dwarf surface, leading to a thermonuclear runaway and the excavation and ejection of material from the O-Ne-Mg region of the white dwarf, whose thin surface layer still has C/O $>$ 1.
Of the five known hydrogen-deticient planetary nebulae 330. A558. 778. IRAS 15154-5258 and IRAS 18333-2357). we have now carried out detailed ORL/CEL abundance studies of two of them.
Of the five known hydrogen-deficient planetary nebulae 30, 58, 78, IRAS 15154-5258 and IRAS 18333-2357), we have now carried out detailed ORL/CEL abundance studies of two of them.
In both cases. we find that the VLTP born-again scenario commonly invoked to account for the production of H-deficient material within an old planetary nebula. and which predicts C/OI. cannot account for the abundance ratios observed. while the collimated outflows seen in both objects seem inconsistent with a single star at the centre.
In both cases, we find that the VLTP born-again scenario commonly invoked to account for the production of H-deficient material within an old planetary nebula, and which predicts $>$ 1, cannot account for the abundance ratios observed, while the collimated outflows seen in both objects seem inconsistent with a single star at the centre.
The presence of knots of cold hydrogen-deficient material has commonly been invoked to account for the observed discrepancy in planetary nebulae whereby ORLs give much higher chemical abundances than CELs for heavy elements (e.g. Liu et al.
The presence of knots of cold hydrogen-deficient material has commonly been invoked to account for the observed discrepancy in planetary nebulae whereby ORLs give much higher chemical abundances than CELs for heavy elements (e.g. Liu et al.
2000. Tsamis et al.
2000, Tsamis et al.
2004. Wesson et al.
2004, Wesson et al.
2005).
2005).
The origin of this H- material in normal nebulae is as yet unknown.
The origin of this H-deficient material in normal nebulae is as yet unknown.
Given the uncertainty at the moment about how the knots in 330 and 558 have been produced. it is difficult to say if the proposed knots present in ‘normal’ nebulae could have a similar origin.
Given the uncertainty at the moment about how the knots in 30 and 58 have been produced, it is difficult to say if the proposed knots present in `normal' nebulae could have a similar origin.
However. we note that it is difficult for the evolution of a single star to explain the morphology and abundances in A330 and 558. and that recent work has suggested that most or all central stars of planetary nebulae could be binary systems (Moe and De Marco 2006).
However, we note that it is difficult for the evolution of a single star to explain the morphology and abundances in 30 and 58, and that recent work has suggested that most or all central stars of planetary nebulae could be binary systems (Moe and De Marco 2006).
However. spatially resolved spectroscopy and high resolution imaging of nebulae with high abundance discrepancy factors. such as NGC 6153. seem to suggest that knots in "normal, nebulae must be very small. «60 AAU across. numerous. and have a smooth distribution. as no clumps are seen in HST images. nor spikes in long slit spectra. but rather a smooth decline of adf from centre to edge. for example in NGC 6153 (Liu et al.
However, spatially resolved spectroscopy and high resolution imaging of nebulae with high abundance discrepancy factors, such as NGC 6153, seem to suggest that knots in `normal' nebulae must be very small, $<$ AU across, numerous, and have a smooth distribution, as no clumps are seen in HST images, nor spikes in long slit spectra, but rather a smooth decline of adf from centre to edge, for example in NGC 6153 (Liu et al.
2000).
2000).
This is in contrast to the highly clumpy distribution of H-deficient material seen in Abell 30. Abell 58 and Abell 78.
This is in contrast to the highly clumpy distribution of H-deficient material seen in Abell 30, Abell 58 and Abell 78.
This work is partly supported by a joint research grant co-sponsored by the Natural Science Foundation of China and the UK's Royal Society.
This work is partly supported by a joint research grant co-sponsored by the Natural Science Foundation of China and the UK's Royal Society.
OD acknowledges funding from NSF-AST-0607111.
OD acknowledges funding from NSF-AST-0607111.
We thank the anonymous referee for a thorough and useful report.
We thank the anonymous referee for a thorough and useful report.
The Galactic diffuse 5-rav. emission is believed to be mostly produced in interactions of cosmic ravs with the matter aud (he radiation fields in the Galaxy. the main production mechanisms being electron non-thermal Dremsstrahlung. Inverse Compton scatterings off the radiation fields and pion decay processes in inelastie collisions of nuclei and matter.
The Galactic diffuse $\gamma$ -ray emission is believed to be mostly produced in interactions of cosmic rays with the matter and the radiation fields in the Galaxy, the main production mechanisms being electron non-thermal Bremsstrahlung, Inverse Compton scatterings off the radiation fields and pion decay processes in inelastic collisions of nuclei and matter.
Although the standard production mechanisms of 5-ravs (Bertschetal.1993) explain generally. well the spatial and energy distribution of the emission below 1 GeV. the model does not match EGRET observations of the 5-rav skv above 1 GeV (IIunteretal.1997)..
Although the standard production mechanisms of $\gamma$ -rays \citep{Bertsch:1993} explain generally well the spatial and energy distribution of the emission below 1 GeV, the model does not match EGRET observations of the $\gamma$ -ray sky above 1 GeV \citep{Hunter:1997we}.
Many. possible explanations have been proposed (o account for the GeV excess
Many possible explanations have been proposed to account for the GeV excess
Observations of the field were performed in April 1999 with the first Unit Telescope (Anti) of the ESOVLT located at the Paranal Observatory.
Observations of the field were performed in April 1999 with the first Unit Telescope (Antu) of the ESO located at the Paranal Observatory.
Lnages were obtained using the FOcal Reducer aud Spectrograph #11 (FOBRS1) instrument. a four-port 2048x CCD detector which can be used both as a high/low resolution spectrograph and an imaging camera.
Images were obtained using the FOcal Reducer and Spectrograph 1 (FORS1) instrument, a four-port $2048 \times 2048 $ CCD detector which can be used both as a high/low resolution spectrograph and an imaging camera.
The instrument was operated in imaging mode with a 1x binning and at its standard angular resolution of 0.2 arcsec/pixel. with a corresponding field of view of 658x678.
The instrument was operated in imaging mode with a $1\times1$ binning and at its standard angular resolution of 0.2 arcsec/pixel, with a corresponding field of view of $6\farcm8 \times 6\farcm8$.
Two 300 s exposures were taken in the Bessel lilters AR (A=6570À: AX=1500A)) and I (AX=7680A: AA= 1380A)) under good seeing conditions (<1").
Two 300 s exposures were taken in the Bessel filters $R$ $\lambda=6570$; $\Delta\lambda=1500$ ) and $I$ $\lambda=7680$; $\Delta\lambda=1380$ ) under good seeing conditions $(\le 1''$ ).
The images have been retrieved [rom the ESO public archive and reduced following standard procedures for debiassing and flatfielding.
The images have been retrieved from the ESO public archive and reduced following standard procedures for debiassing and flatfielding.
Since no standard stars were observed. absolute fIux calibration was performed using as a reference a set οἱ secondary stars detected in the same passbands in theNTT images.
Since no standard stars were observed, absolute flux calibration was performed using as a reference a set of secondary stars detected in the same passbands in the images.
Three 20-min exposures of the Vela pulsar field were taken in Il, (A=6588.27À.. AX= 74.31A)) on 1999 April 4 with the Wide Field Imager (WFI) at the ESO/AIPG 2.2 m telescope.
Three 20-min exposures of the Vela pulsar field were taken in $_\alpha$ $\lambda = 6588.27$, $\Delta \lambda=74.31$ ) on 1999 April 4 with the Wide Field Imager (WFI) at the ESO/MPG 2.2 m telescope.
The WFI is a wide fiekl mosaic camera. composed of eight 2048x4096 pixel CCDs. with a scale of 0.238 arcesec/pixel. providing a field of view of 3377x3277.
The WFI is a wide field mosaic camera, composed of eight $2048\times4096$ pixel CCDs, with a scale of $0.238$ arcsec/pixel, providing a field of view of $33\farcm7\times 32\farcm7$.
To compensate for the loss of signal due to the interchip gaps (23788 and 14733 along right ascension and declination. respectively). the second and third exposures were taken with a relative offset of 30" in RA and Dec. The images were taken in fairly good seeing conditions with EWIIMe079,
To compensate for the loss of signal due to the interchip gaps 8 and 3 along right ascension and declination, respectively), the second and third exposures were taken with a relative offset of $30\arcsec$ in RA and Dec. The images were taken in fairly good seeing conditions with $\simeq0\farcs9$.
Data reduction with the usual CCD processing steps of bias subtraction. flat-lielding and trimming of images. was performed. with using the Mosaic CCD reduction package (AISCRED).
Data reduction with the usual CCD processing steps of bias subtraction, flat-fielding and trimming of images, was performed with using the Mosaic CCD reduction package (MSCRED).
Individual exposures have been finally coadcded ancl cleaned ol cosmic-ray hits using the DRIZZLE software (look Fruchter 1997) and the final image has been corrected [or the effects of Iringing that alfects the CCDs in the red part of the spectrum (Z6500 ΑΔ}.
Individual exposures have been finally coadded and cleaned of cosmic-ray hits using the DRIZZLE software (Hook Fruchter 1997) and the final image has been corrected for the effects of fringing that affects the CCDs in the red part of the spectrum $\gtrsim6500$ ).
Of course. the [lux density of the galactic background radiation also depends on the solar location in the sky.
Of course, the flux density of the galactic background radiation also depends on the solar location in the sky.
Over large solid angles away from the galactic plane the Gus density doa is fairly constant and. σα.) is applicable.
Over large solid angles away from the galactic plane the flux density $I_{\rm gal}$ is fairly constant and \ref{eq2}) ) is applicable.
Llowever. near the galactic plane an enhancement is superposecL on the isotropic component ancl displavs a much richer structure in angle (see an interesting analysis of this problem in the paper of Manning and. Dulk (2001).
However, near the galactic plane an enhancement is superposed on the isotropic component and displays a much richer structure in angle (see an interesting analysis of this problem in the paper of Manning and Dulk (2001).
The sky map may be divided into the two flux density regions using galactic coordinates. as it was shown in Joardar ancl Bhattacharva (2007). but in our study we will not account for the structure of galactic background radiation.
The sky map may be divided into the two flux density regions using galactic coordinates, as it was shown in Joardar and Bhattacharya (2007), but in our study we will not account for the structure of galactic background radiation.
We only have restricted. by Eqs. (2)) and(3))
We only have restricted by Eqs. \ref{eq2}) ) \ref{eq3}) )
in this consideration.
in this consideration.
Using the Ravleigh-Jeans Law. it is not difficult to caleulate the brightness temperature of the galactic background radiation at 9-13.8 MllIz.
Using the Rayleigh-Jeans Law, it is not difficult to calculate the brightness temperature of the galactic background radiation at 9-13.8 MHz.
In particular. for the frequency 25 ΜΗ12 it is about 28300 Ix. that comes to agreement with the measurements reported in the works of Ixevmkin (1971) ane Ixonovalenko et al(2007).
In particular, for the frequency 25 MHz it is about 28300 K that comes to agreement with the measurements reported in the works of Krymkin (1971) and Konovalenko et al.(2007).
The beam solid angle of the WIND-spacecraft antenna is Qua=SR/A sterad (assuming. it receives in a iomiogeneous medium of infinite extent).
The beam solid angle of the WIND-spacecraft antenna is $\Omega_{\rm ant}=8\pi/3$ sterad (assuming, it receives in a homogeneous medium of infinite extent).
Consistent with ]xraus. (1967). its elfective area equals clo=AX?ANT. where A is the wavelength of radio emission.
Consistent with Kraus (1967), its effective area equals $A_0=\lambda^2\times 3/8\pi$, where $\lambda$ is the wavelength of radio emission.
As a measure of he telescope performance. we apply the Svstem Equivalent Solar Flux Density (SESEFD) similar to. the ordinary definition of SEED. (Campbell 2002).
As a measure of the telescope performance, we apply the System Equivalent Solar Flux Density (SESFD) similar to the ordinary definition of SEFD (Campbell 2002).
Consider an 1.0 sfu »oint source. and the antenna temperature due to this source is expressed in terms of the Ix/sfu value defined in Eq. (1)).
Consider an 1.0 sfu point source, and the antenna temperature due to this source is expressed in terms of the K/sfu value defined in Eq. \ref{eq1}) ).
The SIESED is the point source IHux density (in sfu) which xoduces the antenna temperature Ly=7... namely For the antenna array of 720 dipoles with the sensitivity equal to 170.000 Ix/sfu and the svstenm noise temperature 130.000 Ix at 13.8 MlIz. the SIZSED is less 1 δα. whereas for the WIND-spacecralt antenna with the sensitivity. 204 Ix/sfu at the same frequeney. the SESED amounts to 637 sfu.
The SESFD is the point source flux density (in sfu) which produces the antenna temperature $T_{\rm A}= T_{sys}$, namely For the antenna array of 720 dipoles with the sensitivity equal to 170,000 K/sfu and the system noise temperature 130,000 K at 13.8 MHz, the SESFD is less 1 sfu, whereas for the WIND-spacecraft antenna with the sensitivity 204 K/sfu at the same frequency, the SESFD amounts to 637 sfu.
Although the brightness temperature of galactic background radiation is almost the same for receiving by either one dipole or the array. but the brightness temperature of solar radio bursts will be directly proportional to the effective area of the receiving antennas.
Although the brightness temperature of galactic background radiation is almost the same for receiving by either one dipole or the array, but the brightness temperature of solar radio bursts will be directly proportional to the effective area of the receiving antennas.
In this comparative analysis we do not account for any losses in the antennas themselves and their feeder systems. as well as the signal-to-noise ratio per beam due to the back end.
In this comparative analysis we do not account for any losses in the antennas themselves and their feeder systems, as well as the signal-to-noise ratio per beam due to the back end.
Nevertheless. this is quite enough to compare the performance of these instruments.
Nevertheless, this is quite enough to compare the performance of these instruments.
Strictly saving about the system noise temperature for the antenna array of 720 dipoles. it should. be also. pointed out an appreciable (in comparison with galactic background radiation) contribution of quiet solar emission.
Strictly saying about the system noise temperature for the antenna array of 720 dipoles, it should be also pointed out an appreciable (in comparison with galactic background radiation) contribution of quiet solar emission.
In these conditions. due to the quiet solar emission. the system noise temperature increases by about in the dependence of frequency.
In these conditions, due to the quiet solar emission, the system noise temperature increases by about in the dependence of frequency.
li gaiould be also mentioned. about the clillicultices of decameter radio observations for any ground-based dipole in the daytime because of high levels of terrestrial interference that are represented in the bottom panels of Figs.
It should be also mentioned about the difficulties of decameter radio observations for any ground-based dipole in the daytime because of high levels of terrestrial interference that are represented in the bottom panels of Figs.
E. and 2..
\ref{fig1} and \ref{fig2}.
Due to the selectivity of the UTT-2 array of 720 clipoles. it has proved to be a most valuable instrument for. the observing of solar radio bursts in such nonsimple conditions of observation.
Due to the selectivity of the UTR-2 array of 720 dipoles, it has proved to be a most valuable instrument for the observing of solar radio bursts in such nonsimple conditions of observation.
There is also one point that it. should. be remembered.
There is also one point that it should be remembered.
In. the frequency. range. 9-14 MlIZ the solar events often happened with the duration on the order of tens seconds ancl loss (Abranin et al.
In the frequency range 9-14 MHz the solar events often happened with the duration on the order of tens seconds and less (Abranin et al.
1979: Mebnik et al.
1979; Mel'nik et al.
2004: Alelnik et al.
2004; Mel'nik et al.
2005a:b: Chernov ct al.
2005a;b; Chernov et al.
2007: Ixonovalenko et al.
2007; Konovalenko et al.
2007). ancl therefore the temporal resolution of the receiver RAD? of the spacecraft. WIND in 16 s is too rough for their analysis.
2007), and therefore the temporal resolution of the receiver RAD2 of the spacecraft WIND in 16 s is too rough for their analysis.
Sum up. it should. be pointed. out some advantages: ancl disadvantages of space observations of cdecameter racio emission from the solar corona in comparison with erouncl-based. measurements.
Sum up, it should be pointed out some advantages and disadvantages of space observations of decameter radio emission from the solar corona in comparison with ground-based measurements.
The evident advantage is that the space observations permit ones to carry out them continuous in time and ονομά the transpareney bounds of ionosphere (lower 7-9 MlIIz) that. ijs. problematic [or &rouncd-based instruments.
The evident advantage is that the space observations permit ones to carry out them continuous in time and beyond the transparency bounds of ionosphere (lower 7-9 MHz) that is problematic for ground-based instruments.
Although in the former case one can hope that if decameter radiotelescopes with a sullicient. effective area were as many as possible. and they were situated in dillerent places of the Earth. then it would be possible to organize a consistent patrol of solar events to be passed. around. the observatories.
Although in the former case one can hope that if decameter radiotelescopes with a sufficient effective area were as many as possible, and they were situated in different places of the Earth, then it would be possible to organize a consistent patrol of solar events to be passed around the observatories.
At that time such a method. of observations would become almost continuous in time.
At that time such a method of observations would become almost continuous in time.
As weak points of the spacecraft observations. it should be noticed. problems to launch a sullicient elfective antenna for the decameter wavelength range. but any second-rate antenna distinctly reduces to the benefits given by the space observatory due to its space location.
As weak points of the spacecraft observations, it should be noticed problems to launch a sufficient effective antenna for the decameter wavelength range, but any second-rate antenna distinctly reduces to the benefits given by the space observatory due to its space location.
As a result. the sullicient strong events can be detected. but any fine structure is washed out.
As a result, the sufficient strong events can be detected, but any fine structure is washed out.
But it is difficult to have a clear concept about any solar event. without accurate records.
But it is difficult to have a clear concept about any solar event without accurate records.
In this context. the erounc-basecl radiotelescope of ἵνρο UTIt-2 has a clear advantage.
In this context the ground-based radiotelescope of type UTR-2 has a clear advantage.
Now a new large decameter radiotelescope. known as LOPAR (LOw Frequency Altray). is being built in Holland.
Now a new large decameter radiotelescope, known as LOFAR (LOw Frequency ARray), is being built in Holland.
Hs appearance is verv important.
Its appearance is very important.
The point is that in the world. there are a number of small antennas (consisting of one or some dipoles) for the decameter radio astronomy purpose such as the antenna BIRS of Evickson in Tasmania (Lrickson 1997). antenna GBSRBS in NIULAO. (httpi//Zwww.astro.umd.edu/white/gb/). the Culgoora Racliospectrograph (Prestage et al.
The point is that in the world there are a number of small antennas (consisting of one or some dipoles) for the decameter radio astronomy purpose such as the antenna BIRS of Erickson in Tasmania (Erickson 1997), antenna GBSRBS in NRAO $\sim$ white/gb/), the Culgoora Radiospectrograph (Prestage et al.
1994). the antenna of IZMIURAN at 25 MIIZ in Russia (http://helios.
1994), the antenna of IZMIRAN at 25 MHz in Russia (http://helios.